Lawrence Berkeley National Laboratory
Recent Work
Title
RAPIDLY RISING TRANSIENTS in the SUPERNOVA - SUPERLUMINOUS SUPERNOVA GAP
Permalink
https://escholarship.org/uc/item/3738w56r
Journal
Astrophysical Journal, 819(1)
ISSN
0004-637X
Authors
Arcavi, I
Wolf, WM
Howell, DA
et al.
Publication Date
2016-03-01
DOI
10.3847/0004-637X/819/1/35
Peer reviewed
eScholarship.org
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RAPIDLY RISING TRANSIENTS IN THE SUPERNOVA - SUPERLUMINOUS SUPERNOVA GAP
Iair Arcavi1,2, William M. Wolf3, D. Andrew Howell1,3, Lars Bildsten2,3, Giorgos Leloudas4,5, Delphine Hardin6, Szymon Prajs7, Daniel A. Perley5, Gilad Svirski8, Avishay Gal-Yam4, Boaz Katz4, Curtis McCully2,3, S. Bradley Cenko9,10, Chris Lidman11, Mark Sullivan7, Stefano Valenti2,3, Pierre Astier6,
Cristophe Balland6, Ray G. Carlberg13, Alex Conley14, Dominique Fouchez15, Julien Guy6, Reynald Pain6, Nathalie Palanque-Delabrouille16, Kathy Perrett17, Chris J. Pritchet18, Nicolas Regnault6, James Rich16 and
Vanina Ruhlmann-Kleider16
1Las Cumbres Observatory Global Telescope, 6740 Cortona Dr, Suite 102, Goleta, CA 93111, USA [email protected] 2Kavli Institute for Theoretical Physics, University of California, Santa Barbara, CA 93106, USA
3Department of Physics, University of California, Santa Barbara, CA 93106, USA
4Department of Particle Physics and Astrophysics, The Weizmann Institute of Science, Rehovot, 76100, Israel
5Dark Cosmology Centre, Niels Bohr Institute, University of Copenhagen, Juliane Maries Vej 30, DK-2100 Copenhagen, Denmark 6LPNHE, CNRS-IN2P3 and University of Paris VI & VII, F-75005 Paris, France
7School of Physics and Astronomy, University of Southampton, Southampton, SO17 1BJ, UK 8Racah Institute for Physics, The Hebrew University, Jerusalem 91904, Israel
9Astrophysics Science Division, NASA Goddard Space Flight Center, Mail Code 661, Greenbelt, MD 20771, USA 10Joint Space-Science Institute, University of Maryland, College Park, MD 20742, USA
11Australian Astronomical Observatory, PO Box 915, North Ryde, NSW 1670, Australia 12University of Paris-Sud, Orsay, F-91405, France
13Department of Astronomy and Astrophysics, University of Toronto, 50 St. George Street, Toronto, ON M5S 3H8, Canada 14Center for Astrophysics and Space Astronomy, University of Colorado, 389 UCB, Boulder, CO 80309-389, USA
15Aix Marseille Universit´e, CNRS/IN2P3, CPPM UMR 7346, 13288, Marseille, France 16DSM/IRFU/SPP, CEA-Saclay, F-91191 Gif-sur-Yvette, France
17DRDC Ottawa, 3701 Carling Avenue, Ottawa, ON K1A 0Z4, Canada and
18Department of Physics and Astronomy, University of Victoria, P.O. Box 3055, Victoria, BC V8W 3P6, Canada Draft version January 18, 2016
ABSTRACT
We present observations of four rapidly rising (trise ≈ 10 d) transients with peak luminosities
be-tween those of supernovae (SNe) and superluminous SNe (Mpeak≈ −20) - one discovered and followed
by the Palomar Transient Factory (PTF) and three by the Supernova Legacy Survey (SNLS). The light curves resemble those of SN 2011kl, recently shown to be associated with an ultra-long-duration gamma ray burst (GRB), though no GRB was seen to accompany our SNe. The rapid rise to a lumi-nous peak places these events in a unique part of SN phase space, challenging standard SN emission mechanisms. Spectra of the PTF event formally classify it as a Type II SN due to broad Hα emission, but an unusual absorption feature, which can be interpreted as either high velocity Hα (though deeper than in previously known cases) or Si II (as seen in Type Ia SNe), is also observed. We find that existing models of white dwarf detonations, CSM interaction, shock breakout in a wind (or steeper CSM) and magnetar spindown can not readily explain the observations. We consider the possibility that a “Type 1.5 SN” scenario could be the origin of our events. More detailed models for these kinds of transients and more constraining observations of future such events should help better determine their nature.
Subject headings: supernovae: individual (PTF10iam, SNLS04D4ec, SNLS05D2bk, SNLS06D1hc, Dougie)
1. INTRODUCTION
Supernovae (SNe), the explosive deaths of stars, are observed to occur in a variety of types and a spread of luminosities. Most notable is the division into core collapse SNe (Types Ib/c and II; see Filippenko 1997 for a review), which are associated with the deaths of
massive (M & 8M ) stars, and Type Ia SNe,
associ-ated with the thermonuclear disruptions of white dwarfs (Hoyle & Fowler 1960; Hansen & Wheeler 1969, Nomoto 1982a,b; Nomoto, Thielemann & Yokoi 1984; Branch et al. 1985; see Nugent et al. 2011 for the most direct observational evidence of this association). Recently, a third class of explosions has been identified, superlumi-nous SNe (SLSNe), characterized by their high luminos-ity at peak (e.g. Quimby et al. 2007; Smith et al. 2007; Ofek et al. 2007; Gal-Yam et al. 2009; Pastorello et al. 2010; Quimby et al. 2011; see Gal-Yam 2012 for a
review). These events likely originate in massive stars, but clear progenitor scenarios for SLSNe have yet to be determined.
An open question related to the possible connection be-tween core collapse SNe and SLSNe is whether the appar-ent lack of “intermediate” evappar-ents (i.e. SNe with peak ab-solute magnitudes in the range −19 to −21; Fig. 1) is real or just a selection effect. Arcavi et al. (2014) searched for such events in the spectroscopically confirmed H-rich core collapse sample from the Palomar Transient Factory (PTF; Rau et al. 2009; Law et al. 2009). Three events were found in the centers of non-starforming galaxies, and turned out to be tidal disruptions of stars by super-massive black holes (Arcavi et al. 2014). A fourth event, PTF10iam, was shown to be significantly offset from its host center, was in a starforming galaxy and displayed a faster rise to peak magnitude.
−24 −22 −20 −18 −16 −14 0 5 10 15 20
Peak Absolute Magnitude
No. of Events
SNe Ib/c SNe II SLSNe
Figure 1. Peak magnitudes of core collapse SNe (Li et al. 2011; see also Bazin et al. 2009 and Taylor et al. 2014) and super lu-minous SNe (SLSNe; Gal-Yam 2012). Here all strongly interacting (Type IIn) SNe are excluded. A gap between core collapse SNe and SLSNe is apparent.
Here we investigate the nature of PTF10iam as a SN with peak luminosity in the SN-SLSN “gap”, but with a surprisingly rapid (≈ 10 day) rise to peak. We present three additional events with similarly short rise times and peak luminosities discovered and followed by the
Supernova Legacy Survey (SNLS1; Astier et al. 2006).
These events were identified as part of a broader search for SLSNe in the SNLS dataset, the details of which will be discussed in Wolf et al. (in prep). In brief,
mod-erately luminous events (Mpeak & −19) with available
redshifts that were not initially identified as Type Ia SNe were all checked by eye. The three studied here, SNLS04D4ec, SNLS05D2bk, and SNLS06D1hc all exhib-ited short rise times (. 10 days) and luminous peak mag-nitudes (≈ −20).
Various SNe with rapidly evolving light curves, such as PTF09uj (Ofek et al. 2010), SN 2002bj (Poznanski et al. 2011), SN 2010X (Kasliwal et al. 2011), OGLE-2013-SN-079 (Inserra et al. 2015) and a sample of events discovered by Pan-STARRS 1 (Drout et al. 2014), have
been studied in the past. Taddia et al. (2015) show
that “normal” stripped envelope SNe Ib/c can also have a rapid rise to peak. However, none of the previously studied rapid SNe were as luminous at peak as the sample presented here, with two exceptions. One is the transient
“Dougie” (Vink´o et al. 2015). Dougie was discovered by
the ROTSE-IIIb survey and displays a ≈ 7 day rise to an extremely luminous peak magnitude of ≈ −23 in R-band. Being much more luminous than our events, it may be of a completely different nature. The second known rapidly rising luminous event is SN 2011kl (Greiner et al. 2015). SN 2011kl was recently recovered from the optical afterglow of the ultra-long-duration gamma ray burst (GRB) 111209A (Gendre et al. 2013; Stratta et al. 2013; Levan et al. 2014). Both the GRB properties and the SN properties are not similar to any previous GRB-SN event, but the GRB-SN light curves are similar to those of our sample.
1Not to be confused with the acronym for superluminous super-nova - SLSN
We present the observations of our events (and a new host galaxy spectrum of Dougie, confirming its redshift and thus peak absolute magnitude) in §2 and analyze the light curves and spectra of our events in §3, comparing to those of SN 2011kl. We discuss possible physical mecha-nisms for creating the transients in our sample in §4 and summarize in §5.
2. OBSERVATIONS
PTF10iam was discovered by the Palomar 48-inch Oschin Schmidt telescope (P48) as part of the PTF sur-vey. It was classified as a Type II SN following spectra showing a blue continuum and later broad Hα emission. The three SNLS events (SNLS04D4ec, SNLS05D2bk and SNLS06D1hc) were discovered by the deep survey of the Canada France Hawaii Telescope Legacy Survey
(CFHTLS 2002)2, using the CFHT 3.6-meter telescope.
The events were marked as non-Ia SNe based on their light curves (Sullivan et al. 2006). The discovery infor-mation for all four events is presented in Table 1.
2.1. Photometry
The photometry of PTF10iam was released in
Ar-cavi et al. (2014). Here we recover an additional
pre-explosion non-detection upper limit. For the SNLS events, images in the g, r, i and z bands were obtained as part of the SNLS rolling survey. The SN flux was mea-sured by removing a modeled host contamination, which is taken from reference images after PSF matching (see Astier et al. 2006 and references therein for more de-tails).
The photometry for PTF10iam is presented in the AB system, and for the SNLS events in the Vega system, in Table 2. In Figure 2 we present the light curves in observed filters after correcting for Galactic extinction using the Schlafly & Finkbeiner (2011) maps, extracted
via the NASA Extragalctic Database (NED3).
The PTF10iam spectra presented in Section 2.2 do not display noticable Na I D absorption, indicating that any host extinction for PTF10iam is likely small. Addition-ally, we fit the first spectrum of PTF10iam with a variety of temperatures and extinction values. Assuming no ex-tinction, our best blackbody fit gives a temperature T of 11000 K (see Setion 3.1). We find equally good fits up to T = 13200 K and E(B-V) = 0.19, but the fits signif-icantly worsen with higher extinction values regardless of temperature. The best fit is found for T = 11000 K with E(B-V) = 0.09, which corresponds to an extinction of A = 0.2 mag in the PTF R-band (assuming a Cardelli et al. 1989 extinction law). All of our events show simi-lar blackbody temperatures (Setion 3.1), suggesting that they all suffer comparably low host extinction. We thus neglect host extinction for all of our events.
Distance moduli are calculated from spectroscopic red-shifts of the host galaxies, determined from narrow
spec-tral features (Fig. 4). A cosmological model with
H0 = 70 km s−1Mpc−1, Ωm = 0.3 and ΩΛ = 0.7 is
as-sumed throughout.
We present long-term light curves in observed flux for our events in Figure 3. There is no evidence for addi-tional activity other than the main eruptions.
2 http://cfht.hawaii.edu/Science/CFHTLS/ 3 http://ned.ipac.caltech.edu/
Table 1
Discovery details of our events. Error values denote 1σ uncertainties.
Name RA Dec Redshift Discovery Discovery
(J2000) (J2000) Date Mag PTF10iam 15:45:30.85 +54:02:33.0 0.109 2010 May 22 19.14 ± 0.11 SNLS04D4ec 22:16:29.29 −18:11:04.1 0.593 2004 Jul 9 22.70 ± 0.06 SNLS05D2bk 10:02:13.96 +02:05:55.2 0.699 2005 Jan 15 23.33 ± 0.06 SNLS06D1hc 02:24:48.25 −04:56:03.6 0.555 2006 Nov 14 22.55 ± 0.04 −21 −20 −19 −18 −17 −16 Absolute Magnitude PTF10iam P48 Mould−R P60 r P60 i SNLS04D4ec CFHT g+0.5 CFHT r CFHT i−0.5 CFHT z−1.0 −20 0 20 40 60 80 −21 −20 −19 −18 −17 −16
Rest Frame Days
Absolute Magnitude SNLS05D2bk CFHT g+0.5 CFHT r CFHT i−0.5 CFHT z−1.0 −20 0 20 40 60 80
Rest Frame Days
SNLS06D1hc
CFHT g+0.5 CFHT r CFHT i−0.5 CFHT z−1.0
Figure 2. Light curves (in observed filters) of our four rapidly rising luminous transients. Triangles denote 3σ non-detection upper limits. The solid line represents the PTF10iam light curve for comparison to the SNLS events. Empty circles in the SNLS light curves are the expected magnitudes after K-correction to the PTF Mould-R filter in the PTF10iam rest frame (i.e. these points are the ones that can be directly compared to the solid line). The K-corrections are based on the blackbody fits performed for epochs with 3 or more bands observed within 0.5 days (see text for details). All four events exhibit a rapid . 10 day rise to a luminous (≈ −20 mag) peak. SNLS05d2bk shows a second peak approximately 40 days after the main peak.
The SNLS host-galaxy magnitudes were obtained from the SNLS 5-year imaging data set (Hardin et al, in prep), following the general method described in Kronborg et al. (2010). In short, photometry was performed on deep im-age stacks in the ugriz Megacam filters. The deep stacks are constructed by selecting 60% of the best quality
im-ages. Transmission and seeing cuts (FWHM < 1.1”)
were applied. Because we have fewer exposures in the
u-band than in the other u-bands, less stringent quality cuts are applied to these images. For the Deep-D2 field (rele-vant for SNLS objects with “D2” in their name), we use
the Terapix T0006 D2-u stack4, as it incorporates
COS-MOS (Scoville 2007, Koekemoer 2007) CFHT-Megacam-u data that partially overlap with the D2 field and were
Table 2
Photometric observations (upper limits mark 3σ non-detections). The PTF10iam data were presented also in Arcavi et al. (2014), but are given here again for completeness with the addition of a new pre-explosion upper limit. This table is published in its entirety in the electronic version. A portion is shown here for
guidance regarding its form and content.
Object Telescope Filter MJD Mag Error PTF10iam P48 R 55323.462 > 21.222 PTF10iam P48 R 55324.263 > 21.134 PTF10iam P48 R 55324.306 > 21.216 PTF10iam P48 R 55330.491 > 20.914 PTF10iam P48 R 55338.779 > 20.994 PTF10iam P48 R 55345.472 19.142 0.114 PTF10iam P48 R 55346.267 18.973 0.046 PTF10iam P48 R 55346.311 18.905 0.048 PTF10iam P48 R 55351.3 18.451 0.025 not processed in the SNLS pipeline.
The selected images were co-added using the swarp
v2.17.1 package5to produce a large contiguous 1 square
degree “season”-stack (a season corresponds to the 6 con-secutive months during which the field was observed). These “season” frames are further co-added excluding the season during which the supernova exploded.
The source detection and photometry is performed us-ing SExtractor V2.4.4 (Bertin & Arnouts, 1996) in dou-ble image mode. The detection is made in the i band. Zero points are computed using aperture photometry on a tertiary star catalog described in Regnault et al. (2009).
For PTF10iam and Dougie, host-galaxy magnitudes
are taken from the Sloan Digital Sky Survey (SDSS)6
via DR10 (Ahn et al. 2014). All host-galaxy magnitudes are presented in Table 3.
2.2. Spectroscopy
Spectra of PTF10iam were released in Arcavi et al. (2014), and are presented here in Figure 5. A spectrum of the host galaxy of PTF10iam was obtained by SDSS in 2002 and downloaded via DR10. This spectrum is presented in Figure 4.
No spectra were obtained of SNLS04D4ec and
SNLS05D2bk during outburst, but spectra of their host galaxies were taken after the transients faded signifi-cantly. A spectrum of SNLS06D1hc was obtained about three weeks after explosion, but no discernible SN fea-tures (other than a possible blue continuum) can be seen in the spectrum. All SNLS spectra were taken with the FOcal Reducer and Spectrograph (FORS1 and FORS2;
Appenzeller et al. 1998) mounted on the Very Large
Telescope (VLT). The FORS data were reduced using a
mixture of standard IRAF7 tasks and our own routines
that were specifically written to process MOS data from FORS1 and FORS2. For the SNe, we also derive an er-ror spectrum, which is computed from regions in the 2D sky-subtracted spectrum that are free of objects. The
5http://terapix.iap.fr/soft/swarp/ 6http://www.sdss.org
7IRAF, the Image Reduction and Analysis Facility, is a general purpose software system for the reduction and analysis of astro-nomical data. IRAF is written and supported by the National Optical Astronomy Observatories (NOAO) in Tucson, Arizona
0 500 1000 1500 0 5 10 x 10−17 PTF10iam MJD − 54961 F λ [erg s −1 cm −2 Ang −1 ] 0 200 400 600 800 0 0.5 1 1.5 2 x 10−18 SNLS04D4ec MJD − 52795 F λ [erg s −1 cm −2 Ang −1 ] 0 200 400 600 800 0 1 2 x 10−18 SNLS05D2bk MJD − 53328 F λ [erg s −1 cm −2 Ang −1 ] 0 200 400 600 800 0 1 2 3 4 x 10−18 SNLS06D1hc MJD − 53586 F λ [erg s −1 cm −2 Ang −1 ]
Figure 3. Long-term light curves of our events from the PTF and SNLS surveys. No activity is detected outside the main outburst for each event.
spectra are presented in Figure 4.
To verify the redshift (and thus peak luminosity) of
Dougie reported by Vink´o et al. (2015), we observed
its host galaxy with the Low Resolution Imaging Spec-trometer (LRIS; Oke et al. 1995) mounted on the Keck I 10-meter telescope, and with the Gemini Multi-Object Spectrograph (GMOS; Hook et al. 2004) mounted on the Gemini-North 8.1-meter telescope. The LRIS data were reduced using standard IRAF and IDL routines. The Gemini data were reduced using the Gemini IRAF package in addition to custom Python spectral reduction scripts. The highest signal to noise was obtained for the blue part of the LRIS spectrum, and for the red part of the GMOS spectrum, and we present these in Figure 4. We identify several narrow features in the spectrum and use them to determine a redshift of 0.194 (only slightly
different than the value of 0.191 measured by Vink´o et
tem-Table 3
Host galaxy magnitudes for our events and for Dougie. The data for the host galaxies of PTF10iam and Dougie are taken from SDSS. For the SNLS events the magnitudes are taken from deep co-add pre-explosion SNLS images and are given in the Vega system.
Object u g r i z PTF10iam Host 18.76 ± 0.025 17.67 ± 0.006 17.22 ± 0.005 16.90 ± 0.006 16.80 ± 0.015 SNLS04D4ec Host 22.54 ± 0.046 22.90 ± 0.027 21.98 ± 0.022 21.36 ± 0.020 21.18 ± 0.038 SNLS05D2bk Host 22.85 ± 0.057 22.89 ± 0.033 22.03 ± 0.026 21.17 ± 0.021 20.86 ± 0.033 SNLS06D1hc Host 23.56 ± 0.122 23.91 ± 0.067 22.89 ± 0.045 22.30 ± 0.042 22.07 ± 0.078 Dougie Host 22.90 ± 0.481 21.17 ± 0.048 19.88 ± 0.023 19.45 ± 0.024 19.12 ± 0.058 Table 4
Spectroscopic observations. Phases are denoted in rest-frame days relative to peak luminosity. The spectra of PTF10iam were
presented also in Arcavi et al. (2014), but are noted here for completeness.
Object UT Date Phase Telescope [days] (Instrument) PTF10iam 2010 June 8 2 Keck I (LRIS) PTF10iam 2010 July 7 28 Keck I (LRIS) PTF10iam 2010 July 18 38 P200 (DBSP) SNLS04D4ec Host 2007 Jun 19 VLT (FORS1) SNLS05D2bk Host 2007 Jan 28 VLT (FORS2) SNLS06D1hc 2006 Nov 25 5 VLT (FORS1) Dougie Host 2014 Dec 18 Keck I (LRIS) Dougie Host 2015 Feb 19 Gemini N (GMOS) plates).
We thus confirm the peak absolute magnitude of
Dougie (MR ≈ −23) reported by Vink´o et al. (2015),
placing it amongst the brightest SLSNe. It does not fit in the SN-SLSN luminosity gap which is the focus of our interest here. In addition, our observations of Dougie’s host galaxy reveal it to be absorption-feature dominated, unlike the emission-rich hosts of our sample (Fig. 4). Its redder color also indicates a much lower star formation rate (SFR) compared to the other hosts (Table 8). With its extreme luminosity and passive host galaxy, Dougie is likely a different type of event compared to those in our sample and we do not discuss it further in this work.
Our full spectral log is presented in Table 4.
Digi-tal versions of our spectra are available online through the Weizmann Interactive Supernova data REPository
(WISeREP8; Yaron & Gal-Yam 2012).
3. ANALYSIS
3.1. Blackbody Fits and Bolometric Light Curves
We compare each spectrum of PTF10iam to the sum of a blackbody and a scaled host galaxy spectrum, and fit for the blackbody parameters and host scaling factor simultaneously. For the SNLS events we use the SED from epochs when photometric data is available for at least three different bands within half a day (while using each photometric data point only once). We present the SED fits in Figure 6. Our best fit blackbody tempera-tures and radii for all events are presented in the top and middle panels of Figure 7 and in Table 5. All events show very similar temperatures and radii and evolve at similar rates. The temperatures are between those of the Type
8http://wiserep.weizmann.ac.il
Table 5
Best fit blackbody temperatures and radii, and resulting bolometric luminosities, to the spectra of PTF10iam and to the multi-band photometry of the SNLS events. Phases are listed in rest frame days from peak. Errors denote 1σ confidence intervals
for the blackbody fits, and are propagated to the calculated luminosities.
Object Phase T R Lbol
[d] [K] [1015cm] [1043erg s−1] PTF10iam 1.6 11000+800−500 2.68+0.26−0.35 7.51+2.43−1.27 PTF10iam 27.8 < 5600 > 14.24 < 14.22 PTF10iam 37.7 < 5800 > 6.82 < 3.75 SNLS04D4ec 1.2 8600+400−400 3.40+0.34−0.30 4.52+0.90−0.78 SNLS04D4ec 4.4 8700+200−200 3.07 +0.15 −0.14 3.86 +0.37 −0.34 SNLS04D4ec 8.2 8200+300−300 2.90+0.22−0.20 2.71+0.42−0.38 SNLS05D2bk 0.0 11000+300−200 2.41+0.09−0.13 6.06+0.69−0.43 SNLS05D2bk 2.9 9800+300−300 2.85 +0.18 −0.17 5.35 +0.69 −0.63 SNLS05D2bk 13.5 7700+400−400 3.25+0.37−0.31 2.64+0.59−0.51 SNLS05D2bk 15.3 6800+600−500 4.28 +0.71 −0.67 2.80 +1.13 −0.74 SNLS05D2bk 35.8 7100+500−400 3.87+0.48−0.49 2.71+0.85−0.56 SNLS05D2bk 45.8 7200+1200−1000 2.95 +1.03 −0.78 1.66 +1.42 −0.75 SNLS05D2bk 52.3 7200+900−800 3.00+0.80−0.63 1.72+1.04−0.65 SNLS06D1hc −1.9 11100+200−200 2.12+0.08−0.07 4.87+0.36−0.34 SNLS06D1hc 0.1 10600+200−300 2.33 +0.14 −0.09 4.90 +0.38 −0.53 SNLS06D1hc 2.6 10000+200−100 2.57+0.05−0.10 4.72+0.39−0.19 SNLS06D1hc 4.6 9600+200−200 2.57 +0.11 −0.10 4.01 +0.34 −0.32 SNLS06D1hc 6.4 9300+200−200 2.54+0.11−0.11 3.45+0.31−0.29 SNLS06D1hc 16.1 7200+300−300 3.11 +0.28 −0.24 1.85 +0.33 −0.29 SNLS06D1hc 19.2 7000+700−500 3.27+0.52−0.57 1.83+0.85−0.47 SNLS06D1hc 21.2 6600+500−400 3.21 +0.43 −0.44 1.39 +0.47 −0.31 SNLS06D1hc 23.1 7000+2000−1300 2.74+1.39−1.08 1.28+2.22−0.72 SNLS06D1hc 25.0 5800+700−500 4.24 +0.84 −0.86 1.45 +0.84 −0.44
IIb SN 1993J (data from Richmond et al. 1994) and the Type II SN 1998S (data from Fassia et al. 2000), but the radii of our events are larger. PTF10iam appears to display higher blackbody radii than the SNLS events, but this may be due to the fact that it is the only object for which a spectrum was used in the fit, and not cali-brated photometry. The accurate flux calibration of the PTF10iam spectra can not be tested since no multi-color photometry are available for that object.
We use the blackbody radius and temperatures to cal-culate bolometric light curves, and present these light curves in the bottom panel of Figure 7 and in Table 5. We take the brightest bolometric point to be the peak bolometric luminosity for each event.
We use the SNLS blackbody fits to K-correct the SNLS light curves to the PTF10iam Mould R-band observed
2500 3000 3500 4000 4500 5000 5500 6000 6500 7000 7500 O II Ca II O IIIMg Na H PTF10iam Host SNLS04D4ec Host SNLS05D2bk Host SNLS06D1hc Host (and SN) Dougie Host
Rest Wavelength [Ang]
Normalized F
λ
+ Constant
Figure 4. Spectra of the host galaxies of our events and of Dougie. Several narrow emission and/or absorption features are used to securely determine the redshift of each galaxy. The host galaxies of our events all exhibit strong emission features, indicative of ongoing star formation. Dougie’s host, in contrast, shows only absorption features. It is therefore less likely that Dougie was the explosion of a massive star.
frame, and present the results as empty red circles in Figure 2 (only for epochs in the SNLS light curves when blackbody fits are available). These magnitudes can then be compared to the observed Mould-R photometry of PTF10iam (plotted as a solid red line in each SNLS light curve plot). The SNLS and PTF10iam light curves are very similar at the epochs where a comparison is possible. SNLS05D2bk, however, displays a second peak
approx-imately 40 rest-frame days after the main peak. The
second peak is broader and fainter than the main peak but is clearly seen in all filters. Double peaks are com-mon in the rest-frame IR light curves of Type Ia SNe, and have also been seen in the optical bands of Type IIb SNe (intermediate cases between H-rich SNe II and He-dominated SNe Ib) and in at least one energetic Type Ic (H- and He-stripped) SN. For all cases, however, the sec-ond peak is never seen as late as 40 days after the first one (Richmond et al. 1994; Kasen 2006; Arcavi et al. 2011; Kumar et al. 2013; Bufano et al. 2014; Morales-Garoffolo 2014; Nakar et al. 2015; Arcavi et al. in prep).
3.2. Rise Times and Peak Magnitudes
For PTF10iam we calculate the time and magnitude of
peak brightness (tpeakand Mpeakrespectively) by fitting
a 2nd order polynomial to the R-band light curve (in
flux space) around peak. The explosion time (texp) is
taken as the time of zero flux from the polynomial fit. For the SNLS events, which have less densely sampled rises and peaks, but tighter non-detection constraints, we conservatively take the explosion time to be that of
the last non-detection before the first detection. We take the peak time and magnitude to be that of the brightest measured photometric point in z-band for SNLS04D4ec, and in i-band for the other SNLS events. The measured parameters are presented in Table 6.
The rise time of each event (trise) is taken to be the
rest-frame difference between the time of (main) peak and the explosion time (for the SNLS events this is in-terpreted as an upper limit on the rise time). We plot the resulting rise times and peak magnitudes in Figure 8
together with comparison SNe9.
The rise time and peak magnitudes of the compari-son sample were taken from the referenced sources, if stated there explicitly, otherwise they were extracted in the same way as for PTF10iam. Our events have shorter rise times compared to the “standard” SNe shown and similar rise times as the rapidly evolving SNe 2010X, 2002bj, PTF09uj and those in the Drout et al. (2014) 9The following comparison data is used in Figure 8: the SLSNe PTF09cnd from Quimby et al. (2011), SN 2006gy from Smith et al. (2007) and Ofek et al. (2007) and SN 2007bi from Gal-Yam et al. (2009), the normal type Ia SN 2011fe from Vink´o et al. (2012), the average type Ib/c light curve from Drout et al. (2011) and the Type Ib/c r-band sample of Taddia et al. (2015; grey points), the Type IIb SN 2011dh from Arcavi et al. (2011), the peculiar Type IIn SN 1998S from Li et al. (1998), Nakamura et al. (1998), Bignotti et al. (1998), Leonard et al. (2000), Liu et al. (2000) and Fassia et al. (2000), PTF09uj from Ofek et al. (2010), the rapidly evolving SNe 2002bj from Poznanski et al. (2011) and 2010X from Kasliwal et al. (2011), the Pan-STARRS rapidly evolving “gold” sample from Drout et al. (2014; red diamonds), Dougie from Vink´o et al. (2015) and SN 2011kl from Greiner et al. (2015)
3000 4000 5000 6000 7000 8000 9000 Rest Wavelength [Ang]
Normalized F λ + Constant 12d 38d 48d PTF10iam
Figure 5. Spectra of PTF10iam. Phases are shown in rest-frame days from explosion. The first spectrum is mostly a blue contin-uum, the second spectrum displays broad Hα in emission as well as a notable absorption feature bluewards of Hα (marked with an ar-row). The last spectrum is heavily host-galaxy contaminated. The narrow emission and absorption features are from the host galaxy. sample. Our events, however, are much more luminous and are the only ones located in the SN - SLSN gap, aside from SN 2011kl. We exclude from this plot strongly inter-acting Type IIn SNe, identified by narrow emission lines in their spectra (a comparison to such events is shown in Figure 17).
We estimate the ejecta mass corresponding to a par-ticular light curve rise time using the following expres-sion from Wheeler et al. (2014), which follows from Ar-nett (1982) and assumes central deposition of the power source: Mej≈ 0.77M κ 0.1 cm2g−1 −1 v ph 109cm s−1 trise 10 d 2 (1)
Here, vph is the velocity at the photosphere (sometimes
expressed in terms of the kinetic energy of the ejecta).
This scaling, normalized to vph= 109cm s−1 is depicted
in the upper axis of Figure 8.
We translate peak luminosity and rise time to nickel mass (for fully nickel-powered light curves) following Stritzinger & Leibundgut (2005):
MN i M = Lpeak erg s−1/ 6.45 · 1043exp −trise 8.8 d + 1.45 · 1043exp − trise 111.3 d (2) This relation is depicted in the gray dashed lines in
Fig-0 1 2 3 4
Rest Wavelength [Ang] +1.2 +4.4 +8.2 SNLS04D4ec λL λ [10 43 erg s −1 ] 0 1 2 3 4 5
Rest Wavelength [Ang]
−0.0 +2.9 +13.5 +15.3 +35.8 +45.8 +52.3 SNLS05D2bk λL λ [10 43 erg s −1 ] 1000 2000 3000 4000 5000 6000 7000 8000 0 1 2 3 4 −1.9 +0.1 +2.6 +4.6 +6.4 +16.1 +19.2 +21.2 +23.1 +25.0 SNLS06D1hc
Rest Wavelength [Ang]
λL λ [10 43 erg s −1 ]
Figure 6. Blackbody fits to the SEDs of the SNLS events (for epochs when at least three filters were observed within half a day). Epochs are shown in rest frame days relative to peak. The best-fit temperatures and radii are presented in Table 5 and plotted in Figure 7.
Table 6
Light curve parameters for our events (the explosion and peak dates are in the observed frame, while rise times are in the rest frame). Peak magnitudes refer to R-band for PTF10iam, z-band for SNSL04D4ec and i-band for SNLS05D2bk and SNLS06D1hc. Bolometric luminosities are based on blackbody fits (to a spectrum of PTF10iam and to multi-band photometry of the SNLS events). Errors and
confidence bounds denote 1σ uncertainties.
Object texp tpeak Mpeak trise te Peak Lbol
[MJD] [MJD] [days] [days] [1043erg s−1]
PTF10iam 55342.24 ± 0.14 55353.38 ± 0.06 −20.16 ± 0.01 10.05 ± 0.15 2.53 ± 0.38 7.51+2.43−1.28 SNLS04D4ec > 53180.60 53196.58 −20.33 ± 0.06 < 10.03 1.95 ± 0.88 4.52+0.90−0.78 SNLS05D2bk > 53375.58 53385.55 −20.34 ± 0.02 < 5.87 2.76 ± 1.79 6.06+0.69−0.43 SNLS06D1hc > 54039.34 54056.36 −20.22 ± 0.03 < 10.95 3.62 ± 2.12 4.90+0.38−0.53
ure 8 for a few selected nickel masses.
We note that equations (1) and (2) make several sim-plifying assumptions, such as the opacity being constant and the nickel being concentrated in the center, as well
as the photospheric velocity vph being indicative of the
scaling velocity in the model, and the rise to peak being indicative of the effective timescale (i.e. the geometric mean of the diffusion time and the hydrodynamical time) in the model.
Nevertheless, we use these ejecta and nickel mass esti-mates to roughly sketch out the phase space in Figure 8 for which nickel decay can not be the only power source (i.e. the required nickel mass would be larger than the entire ejecta mass). This is the area of the plot to the top
and left of the dashed black line (denoted MN i> Mej).
Since the estimated masses are approximate, this line should be considered a fuzzy limit rather than a strict one. Our events (as well as PTF09uj and the Drout et al. 2014 sample) require very high ratios of nickel mass to total ejected mass (much higher than seen for nickel powered normal Type Ia SNe). Such ratios are seen in models of pure detonations (Sim et al. 2010) and dou-ble detonations (Kromer et al. 2010) of high-mass white dwarfs. We consider these white dwarf detonations in more detail and investigate other possible power sources in Section 4.
3.3. Comparison to SN 2011kl
To the best of our knowledge, SN 2011kl is the only event which has a similar rise time and peak magnitude as our events (Fig. 8). We compare the full light curves of our events to those of SN 2011kl (after removal of its afterglow and host galaxy components; Greiner et al.
2015). SN 2011kl is at a similar redshift (z = 0.677)
as our SNLS events, so the same bands can be roughly
compared to each other. We plot this comparison in
Figure 9. All events show similar peak magnitudes and post-peak decline rates, and some also have compara-ble rise times as SN 2011kl. Here we allow the explo-sion time (set as time zero in Fig. 9) to shift by 2 days for SNLS04D4ec and by 3 days for SNLS06D1hc com-pared to the values derived earlier in order to improve the match with SN 2011kl. These shifts remain consis-tent with the observed upper limits. The resemblance between SNLS06D1hc and SN 2011kl is the most strik-ing.
The similarities in light curve shapes between our events and SN 2011kl indicate that they may all belong the same class. Given the association between SN 2011kl and the ultra-long duration GRB 111209A (Greiner et
al. 2015), we searched for high-energy outbursts consis-tent with the locations and inferred explosions times for our four events. We considered tabulated catalogs for the all-sky InterPlanetary Network (IPN; Hurley et al. 2010), the 8.8 sr Gamma-Ray Burst Monitor on-board
the Fermi satellite (GBM; Meegan et al. 2009), and
the 2 sr Burst Alert Telescope (BAT; Barthelmy et al. 2005) on-board the Swift satellite (Gehrels et al. 2004). No potential counterparts for any event were reported by the Swift -BAT (for which the precise localizations make chance spatial coincidence highly unlikely). While a search of the IPN database revealed several temporal coincidences, the lack of localization provided for most IPN events makes a firm association impossible.
While we find no affirmative evidence for high-energy emission with any of the SNe presented here, we can not rule out an association. The Swift -BAT only observes a modest fraction of the sky at any given moment, so the likelihood of “missing” an associated GRB is quite high. For the IPN, which effectively provides all-sky coverage with a 100% duty cycle, the low sensitivity to very long
duration (∆t > 103s) transients does not guarantee a
detection due to the nature of the triggering algorithm (e.g., Levan et al. 2014). We conclude that present data do not allow us to confirm or refute an association be-tween the SNe presented here and GRB-like transients.
3.4. PTF10iam Spectral Features
The spectrum of PTF10iam obtained 28 (rest frame) days after peak is the first to show significant broad fea-tures. Most notable is broad Hα emission and a broad absorption feature just bluewards of Hα. We present Su-perfit (Howell et al. 2005) results, comparing this spec-trum to that of the Type IIP SN 1999em (from Hamuy et al. 2011) and to the peculiar Type Ia SN 1999ac (from Garavini et al. 2005) in Figure 10.
The spectrum of SN 1999em fits most of the features of PTF10iam quite well, confirming the initial SN II clas-sification of this event. However the Hα P-Cygni profile is very different. The absorption feature in PTF10iam is notably more blueshifted, meaning it could be the re-sult of a high-velocity hydrogen component. To test this possibility, we remove a low order polynomial from the spectrum of PTF10iam and plot the area around Hα in Figure 11. We do the same for an earlier spectrum of SN 1999em (Leonard et al. 2002) and for a spectrum of the Type II SN 2013ej (Valenti et al. 2013). Un-like PTF10iam, these are both SNe IIP (with a plateau in their light curve), but they also show an absorption feature blueshifted from the “main” P-Cygni component
0.4 1 1.5 2
Rest Frame Days Since Explosion T BB [10 4 K] SN 1993J SN 1998S PTF10iam SNLS04D4ec SNLS05D2bk SNLS06D1hc 0.5 1 5 10 15
Rest Frame Days Since Explosion R BB [10 15 cm] 0 10 20 30 40 50 60 1043 1044
Rest Frame Days Since Explosion
Bolometric L [erg s
−1
]
Figure 7. Best fit blackbody temperatures (top) and radii (mid-dle) for our events, and for the Type II SNe 1993J (from Rich-mond et al. 1994; explosion date from Filippenko et al. 1993) and 1998S (from Fassia et al. 2000; explosion date from Chugai 2001) for comparison. The resulting blackbody bolometric lumi-nosities are shown in the bottom panel. Our events show similar blackbody evolution amongst themselves. They have temperatures between those of the partially hydrogen-stripped IIb SN 1993J and the hydrogen-rich SN 1998S, but show more extended blackbody radii than both comparison events.
of Hα. Chugai et al. (2007) interpret a similar feature (appearing at ≈ 50 days post explosion) as high veloc-ity Hα for SN 1999em. Valenti et al. (2013) interpret the early-phase appearance of this feature as Si II for SN 2013ej (see also Parrent et al. 2015 for a discussion on similar features in Type I SNe). It may be possible that this feature is related to Si II at early phases and high velocity Hα at later phases. In either case, however, the feature is much more pronounced in PTF10iam than in the comparison Type II SNe mentioned above.
We fit the full profile from the spectrum of PTF10iam with two Hα P-Cygni components, one at “normal” ve-locities and one at high veve-locities. Each P-Cygni profile is made of two equal-width but shifted and inverted Gaus-sians (the best fit Gaussian parameters are presented in Table 7). The high velocity hydrogen interpretation has the advantage that it fits both the blueshifted absorp-tion feature and the possible high velocity emission tail redwards of Hα, which would not be explained by Si II.
Chugai et al. (2007) consider high velocity hydrogen as a sign of interaction of the SN ejecta with the CSM. Such interaction could also power the light curve of PTF10iam, and would explain its extended blackbody radius.
Table 7
Parameters of the four best fit Gaussian functions used to reproduce the Hα emission profile and bluer absorption feature of
PTF10iam presented in Figure 11. The mean offset is shown in 103km s−1relative to rest-frame Hα, the 1σ width is shown in 103km s−1and the normalization is shown in relative units.
Bounds indicate 67% confidence intervals. Parameter “Normal” Velocity High Velocity
Component Component Emission:
Mean Offset 0 (fixed) 0 (fixed) Width 1.9 (1.5,2.3) 4.8 (4.8,4.9)
Normalization 1 1.11
Absorption:
Mean Offset −2.5 (−1.5,−3.4) −14.5 Width Fixed to the same values as in emission
Normalization 0.68 1.66
A similar spectral feature was seen in a spectrum of SN 1998S (Li et al. 1998), though still weaker and not as blue as that of PTF10iam (Fig. 12). Lentz et al. (2001) model this feature as a blend of Fe II and Si II, and rule out a high velocity Hα origin. We compare the light curves of PTF10iam and SN 1998S and find that SN 1998S was not as fast to rise nor as luminous at peak as PTF10iam, though both events do have similar post-peak decline rates (Fig. 13).
Given that we are not able to find a Type II SN with a similar absorption feature and light curve behavior, we turn to the Si II interpretation. The similarities with
the Type Ia SN 1999ac (Fig. 10) is intriguing10. Not
only does the Si II in the spectrum of SN 1999ac align well with the broad absorption feature of PTF10iam, but many other features (except for the hydrogen lines) fit quite well. This similarity suggests that the spectrum of PTF10iam may be explained as that of a (peculiar) 10SN 1999ac is a peculiar Type Ia of the 99aa-like class (Li et al. 2001, Garavini et al. 2004).
1 10 100 −24 −23 −22 −21 −20 −19 −18 −17 −16 SNe Ib/c PTF09cnd (SLSN−I) SN2006gy (SLSN−II) SN2007bi (SLSN−R) SN 2011fe (Ia) SN Ib/c Template SN 2011dh (SN IIb) trise [Days] Peak R or r Magnitude SN 1998S SN 2010X PTF09uj SN 2002bj PS1 Fast−Evolvers (gold sample) SN 2011kl SNLS04D4ec SNLS05D2bk SNLS06D1hc PTF10iam Dougie 0.01 0.07 0.3 0.8 7 28 77 ~Mej [M⊙κ−1
0 . 1v9] (assuming a central power source and constant opacity)
80 32 13 5 2 0.8 0.3 M N i= 0.1M ⊙ 0.01 0.02 0.05 MNi< Me j MNi >Me j
Figure 8. Peak magnitude vs. rise time of our events (upper limits for the SNLS rise times) compared to other SNe (see text for references). All comparison data peak magnitudes and rise times are in the observed R or r-band. Rise times are in the rest frame of each event. Ejecta mass estimates are normalized to an expansion velocity of 10, 000 km s−1(see text for details, also regarding the calculated nickel masses), and should only be considered approximate. Our events have shorter rise times compared to most SNe, and are more luminous than all similarly-rapid events (except for Dougie, which is a clear outlier in this context). The only similar event to ours is SN 2011kl, which was accompanied by an ultra-long-duration GRB (Greiner et al. 2015). The positions of our events in this phase space require either a very high nickel to ejecta mass ratio or an alternative dominant power source to nickel decay.
Type Ia SN with added broad hydrogen features (see insets in Figure 10, which present the difference between PTF10iam and SN 1999ac around each of the hydrogen lines).
Some Type Ia SNe have been observed to interact with a H-rich CSM (these are known as Type Ia-CSM events or “02ic-likes”; Hamuy et al. 2003; Livio & Riess 2003; Dilday et al. 2012; Silverman et al. 2013; Leloudas et al. 2015a). To test whether PTF10iam could be such an event, we compare its early light curve to those of the normal Type Ia SN 2011fe, the Ia-CSM SNe 2005gj and PTF11kx, and the 91bg-like SN 1999by (Fig. 14). PTF10iam clearly rises more rapidly than SN 2011fe, ruling out an additional interaction component on top of a normal SN Ia, and much faster than the inter-acting SN 2005gj. The rise of PTF10iam is similar to those of SN 1999by and PTF11kx, but PTF10iam de-clines much more slowly than SN 1999by, is more than
2 magnitudes brighter at peak than SN 1999by and is a magnitude brighter than PTF11kx. If PTF10iam were a 91bg-like event with added interaction, or a PTF11kx-like with even stronger interaction, the additional lumi-nosity would require interaction power to dominate the emission. However the spectra of PTF10iam show no signs of interaction. Specifically they do not display the narrow or intermediate-width Hα in their spectra that are prominent in Ia-CSM events.
We conclude that the spectrum of PTF10iam is incon-sistent with a strong interaction power source, but that it may be interpreted as that of either a peculiar Type II SN, or a hybrid Type Ia - Type II event. These inter-pretations will be discussed below in the context of the extreme light curve behavior of PTF10iam.
−24 −23 −22 −21 −20 −19 −18 −17 −16 −15 Absolute Magnitude g+1.5 r i−1.5 z−3 PTF10iam SN 2011kl g+1.5 r i−1.5 z−3 SNLS04D4ec SN 2011kl −20 0 20 40 60 −24 −23 −22 −21 −20 −19 −18 −17 −16 −15
Rest Frame Days
Absolute Magnitude g+1.5 r i−1.5 z−3 SNLS05D2bk SN 2011kl −20 0 20 40 60
Rest Frame Days
g+1.5 r i−1.5 z−3 SNLS06D1hc SN 2011kl
Figure 9. Comparison of the light curves of our events (filled symbols and lines) to SN 2011kl (empty symbols; Greiner et al. 2015), a SN that accompanied an ultra-long-duration GRB. Time zero for SN 2011kl is set to the time of the GRB trigger. For our events it is set to the estimated time of explosion (with an offset of 2 days for SNLS04D4ec and 3 days for SNLS06D1hc, to improve the match). No brightness matching was applied. PTF10iam is at a substantially different redshift than SN 2011kl (z=0.109 vs. z=0.677), so the observed wavelength coverages do not match. SNLS06D1hc, on the other hand, is at a very similar redshift as SN 2011kl, and the two events appear almost identical in their light curve shape (in each filter), indicating that they may be members of the same class of explosions.
3.5.1. Photometric Analysis
We fit the host-galaxy ugriz magnitudes with spec-tral energy distributions (SEDs) computed using PE-GASE2 (Fioc & Rocca-Volmerange 1997, 1999) stellar
population synthesis models. We use the eight
star-forming scenarios described in Table 1 of Le Borgne & Rocca-Volmerange (2002) and the default modeling of internal dust presented therein, together with the initial mass function of Rana & Basu (1992) to compute stellar
masses and recent (averaged over the last 5 · 108yrs)
spe-cific star formation rate (sSFR). Uncertainties are eval-uated through a Monte Carlo propagation of the host-galaxy magnitude uncertainties.
3.5.2. Spectroscopic Analysis
We scale the host galaxy spectra (Fig. 4) to the host galaxy photometry (Table 3) and correct for foreground Galactic extinction (Schlafly & Finkbeiner 2011). The flux of each emission line was measured by fitting a Gaus-sian. We fixed the FWHM of the weakest lines to those
of the lines that were significantly detected.
The host of PTF10iam is the most nearby one of our sample and has the highest signal to noise ratio. This allows for an estimation of the host extinction. Based on the Balmer decrement (Osterbrock 1989), we find E(B-V) = 0.52 ± 0.13. We adopt this redenning for de-riving SFRs, but we caution that it should be considered an upper limit due to the presence of stellar absorption (which affects Hβ more than Hα). This host-integrated extinction does not necessarily affect the line of sight to PTF10iam and may originate in a dusty region be-hind the SN (indeed, we rule out significant extinction for PTF10iam in Section 2.1). After correcting for this extinction, we derive SFRs from the luminosity of the Hα line and (separately) from the luminosity of the [O II] line. For both we use the relations in Kennicutt et al. (1998), corrected to a Chabrier IMF by dividing by a factor of 1.7. Both SFR estimates agree within the uncertainties (Table 8), which contain the measurement error, the host reddening uncertainty and the systematic
Hγ Hβ Hα SN1999em (20d) PTF10iam (28d) Normalized F λ 3000 4000 5000 6000 7000 8000 9000 Hγ Hβ Hα Si II SN1999ac (11d) PTF10iam (28d)
Rest Wavelength [Ang]
Normalized F λ −2 −1 0 1 2 Hα −2 −1 0 1 2 v [104 km/s] Hβ −2 −1 0 1 2 ∆ Fλ Hγ
Figure 10. Superfit comparisons of PTF10iam with the Type II SN 1999em (from Hamuy et al. 2001; top) and the Type Ia SN 1999ac (from Garavini et al. 2005; bottom) with best fit extinction and host-contamination corrections applied. The fit to SN 1999em matches most of the spectral features, except for the absorption feature near 6200 ˚A. The SN Ia fit is able to match this feature as Si II, as well as other features in the spectrum, with the major difference being that PTF10iam has additional broad hydrogen emission lines (insets show the difference between the superfit host and extinction-corrected spectral fluxes of PTF10iam and SN 1999ac around the denoted hydrogen lines; zero flux is marked by the horizontal dotted line; the narrow features are from the not-fully subtracted host of PTF10iam). uncertainties of the conversion relations. We compute
metallicities based on line flux ratios and the calibra-tions presented in Pettini & Pagel (2004) and Kewley & Ellison (2008). The results are presented in Table 8. For metallicities that are based on the R23 scale (Mc-Gaugh 1991, Kobulnicky & Kewley 2004; denoted M91 and KK04 respectively) the upper branch solution is se-lected based on criteria in Kewley & Ellison (2008).
The host galaxies of the SNLS objects are more dis-tant, causing the Hα region to be redshifted outside the observed wavelength. Although the Hβ line is detected in the hosts of SNLS04D4ec and SNLS05D2bk, the non-detection of higher order Balmer lines prevents us from
deriving accurate estimates of the host extinction. A
nominal value of E(B-V) > 1.1 mag can be derived from the upper limit of the Hγ flux. However this derivation is further complicated by signs of stellar absorption that affect the regions of the higher order Balmer lines, and that are difficult to correct for with the signal to noise ratio of these spectra. For SNLS06D1hc we do not detect any Balmer lines.
Because of this uncertainty, we do not apply any host reddening correction to the hosts of the SNLS events. SFRs are calculated based on the luminosity of only the [O II] line, and we compute metallicities only in the R23 scale. The upper branch solution was selected for all the galaxies based on the low [O III]/[O II] ratio (e.g. Na-gao et al. 2006). The uncertainty in host reddening does
−3.5 −2.6 −1.7 −0.7 0.2 1.1 2.0 SN1999em (12d) SN2013ej (25d) PTF10iam (28d) Hα Si II? Hα HV Hα? Velocity [104 km s−1] Normzalied F λ + Constant "Normal" Velocity Hα High Velocity Hα 5800 6000 6200 6400 6600 6800 7000
Rest Wavelength [Ang]
Figure 11. The Hα region in PTF10iam (blue; spectrum taken 28 rest frame days after peak) compared to spectra of SN1999em (Leonard et al. 2002) and SN2013ej (Valenti et al. 2013; red) after removing a low order polynomial from each spectrum. The narrow emission lines in the PTF10iam spectrum are from the host galaxy. The absorption feature at ≈ 6200˚A could be related to high velocity hydrogen, a sign of possible CSM interaction (as interpreted by Chugai et al. 2007 for a later appearance of a similar feature in SN1999em), or to Si II (as interpreted by Valenti et al. 2013 for SN2013ej). In PTF10iam, however, this features is much deeper. We plot the best fit to a sum of four Gaussian functions (black), two representing a “normal” Hα P-Cygni profile and two representing a high velocity P-Cygni profile. The profiles provide a reasonable fit to the features, consistent with the high velocity Hα interpretation, but the absorption depth would be greater than any previously observed high velocity hydrogen feature.
not affect this choice as this ratio would become even lower if a non-negligible extinction is assumed. Due to the non-detection of any Balmer lines in the spectrum of SNLS06D1hc we can not provide any metallicity mea-surements for its host. Our results are present in Table 8 (the Dougie host galaxy does not display any emission lines, and is therefore excluded from this analysis).
The metallicities of all hosts are close to solar or super-solar and the galaxies show clear signs of an evolved stel-lar population, such as stelstel-lar absorption. These galaxies are markedly different from the hosts of H-poor SLSNe
that have been shown to be preferentially star bursting dwarf galaxies (e.g. Neill et al. 2011, Lunnan et al. 2014, Leloudas et al. 2015b).
4. POSSIBLE LIGHT CURVE POWER SOURCES The rapid rise and luminous peak of our events chal-lenge traditional SN power sources. Nickel-decay power would require very high nickel to total mass ratios which are seen in models of pure and double detonations of carbon-oxygen white dwarfs (but not observed in normal Type Ia SNe). We compare our data to models of such detonations and to general energy conservation consid-erations for high nickel mass explosions. We then turn to massive stars and consider three other possible power sources: interaction with the CSM, shock breakout in an optically thick wind and magnetar spindown.
4.1. White Dwarf Detonation
Sim et al. (2010) investigated pure detonations of
sub-Chandrasekhar carbon-oxygen white dwarfs by arti-ficially igniting them in the center. For their most
mas-sive white dwarf (MW D= 1.15M ) they find high nickel
to total mass ratios and consequently rapidly rising lumi-nous light curves. The same behavior is seen by Kromer et al. (2010) who investigate double detonations of white dwarfs (detonating the base of the helium shell, causing a second detonation inside the carbon-oxygen core).
Since hydrogen is not expected to show up in the spec-tra of exploding white dwarfs, we focus on the SNLS events for now. We compare our observed light curves to those modeled by Sim et al. (2010) and Kromer et al. (2010), and find that none match the models well in all filters simultaneously. Given the redshift of our events (z ≈ 0.6), we compare observed rband with model U -band, observed i-band with model B-band and observed z-band with model V -band. We present two of the clos-est matches between the data and the models in Figure 15. As can be seen, even for these cases, the match is unsatisfactory.
An important caveat to these comparisons is that Sim
et al. (2010) and Kromer et al. (2010) do not
con-sider iron group elements in their models. Such elements could introduce additional blue-band opacities, and for the Kromer et al. (2010) models could also influence the nucleosynthesis yields in the helium shell, further affect-ing the light curves.
The poor match of the post-peak light curve between the models and our data disfavor this interpretation for the SNLS events. The hydrogen seen in the spectrum of PTF10iam disfavors this scenario also for that event.
However, there have been suggestions of an explosion channel that would involve white dwarfs detonating in-side hydrogen-rich envelopes - so-called “Type 1.5” SNe (Arnett 1969; Iben & Renzini 1983; Lau et al. 2008). Such SNe are expected to occur when carbon is explo-sively ignited in the core of an intermediate-mass star during its AGB phase. These explosions could be simi-lar to Type Ia SNe with the additiona of hydrogen-rich
ejecta. It is therefore reasonable to expect that they
would show signs of hydrogen in their spectra (coming from the envelope), possibly in addition to deep Si II absorption (as seen in SNe Ia), and would synthesize large amounts of nickel, generating a luminous light curve peak.
3500 4000 4500 5000 5500 6000 6500 7000 7500 SN1998S (36d)
PTF10iam (28d)
Rest Wavelength [Ang]
Normalized F
λ
Figure 12. A comparison of a spectrum of PTF10iam and SN 19998S (from Lentz et al. 2001). The broad absorption feature blueward of Hα interpreted as either high velocity Hα or Si II, is marked in both spectra. The feature is stronger and bluer for PTF10iam.
Table 8
Properties of the host galaxies of our events and of Dougie obtained using fits to host ugriz photometry from Table 3 and analysis of the host emission lines (when available) from the spectra presented in Figure 4. Errors denote 1σ uncertainties. For Dougie we find zero SFR
from the photometric analysis (also when varying the input magnitudes in the Monte Carlo simulation; formally this gives a limit of log (SF R) < −3).
Object Photometric Analysis Spectroscopic Analysis
log (M ) log(sSF R) Hα SFR OII SFR 12 + log(O/H)
[log (M )] [log yr−1] [M yr−1] [M yr−1] (M91) (KK04) (N2) (O3N2) PTF10iam Host 10.40 ± 0.08 −9.70 ± 0.07 4.86 ± 1.58 11.77 ± 7.65 8.67 ± 0.06 8.87 8.57 ± 0.05 8.62 ± 0.04 SNLS04D4ec Host 9.77 ± 0.07 −9.36 ± 0.09 n/a 2.87 ± 0.82 8.59 ± 0.08 8.75 n/a n/a SNLS05D2bk Host 10.27 ± 0.08 −9.46 ± 0.08 n/a 5.06 ± 1.43 8.73 ± 0.03 8.93 n/a n/a SNLS06D1hc Host 9.39 ± 0.09 −9.44 ± 0.09 n/a 0.21 ± 0.07 n/a n/a n/a n/a
Dougie 10.35 ± 0.04 no SF n/a n/a n/a n/a n/a n/a
−20 0 20 40 60 80 −20.5 −20 −19.5 −19 −18.5 −18 −17.5 −17 −16.5
Rest Frame Days Absolute R−Band Magnitude SN1998S
PTF10iam
Figure 13. Light curve comparison between PTF10iam and SN 1998S (data from Li et al. 1998, Nakamura et al. 1998, Bignotti et al. 1998, Leonard et al. 2000, Liu et al. 2000 and Fassia et al. 2000). PTF10iam is faster to rise and more luminous, but the decline rates of both events are very similar.
The progenitors of Type 1.5 SNe are expected to be metal poor, while the derived metallicity for the host galaxies for our events is close to solar (Table 8; though these are global values and not specific to the SN sites). Sparks & Stecher (1974) suggested that a white dwarf spiraling in to the core of a non-degenerate companion and merging with it (the so-called “core-degenerate” sce-nario) could give rise to a similar explosion, without obvi-ous metallicity constraints. In that case, however, some or all of the envelope of the companion might be ejected during the in-spiral. This scenario has thus been used as a possible progenitor channel for Type Ia-CSM events or for events with fast-moving carbon spectral components (Soker et al. 2014).
Since it’s not clear exactly what to expect for true Type 1.5 SNe (assuming they even exist in nature), we now relax many of the assumptions used to derive nickel-powered light curve properties and test only global en-ergy conservation. Katz et al. (2013) present a method for calculating the nickel mass required to power a given bolometric light curve, which relies only on the assump-tions of homologous expansion, radiation-dominated in-ternal energy and nickel decay being the sole power
−50 −40 −30 −20 −10 0 10 20 −20 −19.5 −19 −18.5 −18 −17.5 SN 2011fe −0.86 mag SN 1999by −2.36 mag PTF10iam Polynomial Fit SN 2005gj PTF11kx −1.00 mag
Rest Frame Days
Absolute Magnitude
Figure 14. Early light curve of PTF10iam (red circles are de-tections and the triangle is an upper limit) and the parabolic fit to the pre-peak data used to infer the rise time (solid red line). We compare the rise of PTF10iam to those of the normal Type Ia SN 2011fe (dashed purple line; data from Vinko et al. 2012; dis-tance modulus from Lee & Jang 2012), the 91bg-like SN 1999by (dot-dashed blue line; data from Garnavich et al. 2004; distance modulus from NED), the Ia-CSM PTF11kx (dotted dark green line; Firth et al. 2015) - all shifted in brightness to match the peak of PTF10iam - and the Ia-CSM SN 2005gj (dotted light green line; Aldering et al. 2006). PTF10iam has a faster rise compared to SN 2011fe and SN 2005gj. The rise of PTF10iam, SN ˙1999by and PTF11kx are similar, but their peak magnitudes are very differ-ent. If this difference were due to interaction power it should have imprinted strong CSM signatures in the spectrum.
source. Their argument is that at late enough times
(when the internal energy becomes negligible), the total radiated luminosity equals the total energy deposited by nickel plus the energy lost to expansion. This translates to: Z tlate texp Q (t) ·t dt = Z tlate texp L (t) ·t dt (3)
where Q (t) is the energy deposition from nickel decay,
L (t) is the total radiated luminosity, and tlate is a late
enough time when the internal energy is negligible (Katz
et al. 2013 indicate tlate& 40 days).
We assume full trapping of the positrons, and take the γ-ray optical depth to be:
τγ(t) =
T0
t
−2
(4)
where T0 is left as a free parameter11.
We use the t > 25 days bolometric data points of SNLS05D2bk and SNLS06D1hc from section 3.1 to cal-culate the right hand side of Equation (3), performing a linear interpolation at 0.1 day intervals for the
nu-merical integration, and fit for the nickel mass MN i
and γ escape timescale T0 used to calculate the left
hand side (while keeping the explosion time texp
con-stant at the values listed in Table 6). We find that
for SNLS05D2bk, T0 > 200 days (essentially full
γ-trapping) and MN i= 1.88 ± 0.17M . For SNLS06D1hc
we find T0 = 48+162−15 days (i.e. almost full trapping)
11 Under certain additional assumptions, T0 can be connected to the explosion energy and to the mass and density profile of the ejecta (Clocchiatti & Wheeler 1997)
−21 −20 −19 −18 −17 −16
Rest Frame Days
Absolute Magnitude SNLS04D4ec Model U Model B Model V Data r Data i Data z 0 10 20 30 40 50 −21 −20 −19 −18 −17 −16
Rest Frame Days
Absolute Magnitude
SNLS05D2bk
Figure 15. Two of the closest fits between white dwarf detonation models and our light curves. Top: SNLS04D4ec compared to the Sim et al. (2010) 1.15M model. Bottom: SNLS05D2bk compared to the Kromer et al. (2010) model number 6. Both SNe do not trace the model post-peak declines nor their color evolution. The other models from the Sim et al. (2010) and Kromer et al. (2010) sets were even further from the data, and none were able to reasonably fit SNLS06D1hc at any phase.
and a nickel mass very close to the carbon ignition mass
MN i= 1.38+0.07−0.11M (1σ confidence bounds). The errors
on MN i and T0 were estimated using 500 Monte Carlo
simulated fits. The bolometric light curve, total radiated energy and best fits are plotted in Figure 16.
The integrated bolometric luminosities of
SNLS05D2bk and SNLS06D1hc at late times are consistent with a high-mass nickel decay power source. However, the number of late-time data points for the fit is small and no constraints are provided for the ejecta mass without additional assumptions. While this method does not provide a strong argument for high mass nickel decay indeed being the main power source of the light curves, it is an indication that this possibility is not completely ruled out by the data. More detailed models could perhaps test whether the addition of a hydrogen envelope, as expected for Type 1.5 SNe, can account for the differences in light curves between the detonation models and our observations.
1043 1044 Bolometric L [erg s −1 ] SNLS05D2bk 1043 1044 SNLS06D1hc 0 10 20 30 40 50 60 0 0.5 1 1.5 2 2.5 3 3.5 4 Bolometric ∫ [L(t) ⋅t]dt [10 51 erg * days]
Rest Frame Days Since Explosion 0 10 20 30 40
0 0.5 1 1.5
Rest Frame Days Since Explosion
Figure 16. Fits to the t > 25 days time-integrated bolometric luminosity (bottom) of SNLS05D2bk and SNLS06D1hc with a nickel-powered light curve (dashed lines are for full γ-trapping, visibly different from the solid lines only for SNLS06D1hc). The integrated luminosity is consistent with a nickel-decay power source for the light curves, though there are only few data points and they do not extend to late enough times to make this determination secure. Top plots show the instantaneous luminosity.
is an intriguing possibility, but not a conclusive interpre-tation given the lack of detailed model predictions for this scenario and the lack of strong constraints on the long-term bolometric light curves of our events. We now turn to explosion scenarios involving the core collapse of massive stars.
4.2. CSM Interaction
If a massive star explodes in a dense CSM, the collision of the ejecta with that CSM can produce strong emission and be a major light-curve power source. Depending on the distribution of the CSM, the light curve can be made luminous and either rapidly or slowly evolving. This is the common interpretation of Type IIn SNe (Schlegel 1990), which indeed exhibit luminous light curves with varying rise times (e.g. Kiewe et al. 2012; Ofek et al. 2014b). In addition, if the peculiar absorption feature in the spectrum of PTF10iam is high velocity Hα, then it may be evidence of CSM interaction (Chugai et al. 2007). To compare the peak magnitudes and rise times of our events to those of SNe IIn, we follow Ofek et al.
(2014b) and fit an exponential rise of the form:
L = L0· {1 − exp [(texp− t) /te]} (5)
to each light curve (between discovery and peak). The
free parameter here is te which is treated as a
charac-teristic timescale describing the rise. We plot the best
fit tefor our events, compared to the Ofek et al (2014b)
Type IIn sample in Figure 17.
We find that our events are comparable in peak lu-minosity to the brightest SNe IIn from the Ofek et al. (2014b) sample, but that ours rise faster. Another dis-crepancy with the CSM model is that the spectra of PTF10iam and of SNLS06D1hc do not display the strong intermediate-width and narrow Balmer-series emission
features characteristic of SNe IIn (e.g. Kiewe et al.
2012). Moriya & Tominaga (2012) suggest that a
shal-low density profile (ρCSM∝r−w with w . 1) of the CSM
could allow interaction to power the light curve while not creating IIn-like features in the spectra. However we show below that the light curves of our events im-ply a much steeper density profile of w > 2. Recently, Smith et al. (2015) suggested a new model to account for
100 101 102 1042 1043 1044 te [Days] Peak L bol [erg s −1 ] PTF10iam SNLS04D4ec SNLS05D2bk SNLS06D1hc
SNe IIn (Ofek et al. 2014b)
Figure 17. Estimated peak bolometric luminosities vs. the ex-ponential rise fit parameter teof our events compared to Type IIn SNe from Ofek et al. (2014b). Our events rise more rapidly than interaction-powered SNe.
events with interaction-powered light curves but showing no narrow emission lines in their spectra. The model in-volves a rather complex non-symmetrical CSM distribu-tion, and can explain the observations of SNe 1998S and PTF11iqb as weakly interacting events. Our events are much more luminous, requiring strong interaction, and are thus not readily explained by this model.
In summary, our events have shorter rise times and higher peak luminosities compared to other SNe IIn and they lack IIn-like spectral features (while displaying in-dications of a steep density profile, see below). If our events are powered by CSM interaction, then the initial conditions must be different than for most CSM-powered
SNe. One possibility is brief interaction with a CSM
clump or shell (formed for e.g. in the scenario suggested by Quataert et al. 2015).
4.3. Shock Breakout in a Wind (SBW)
First light from a propagating shock in a SN will emerge when the optical depth τ is approximately equal to c/v (with c the speed of light and v the speed of the SN shock). This is known as the shock breakout (e.g. Colgate 1974, Weaver 1976, and more recently Nakar & Sari 2010, Rabinak & Waxman 2011). The duration of such a signal will be smeared by the light crossing time of the radius at shock breakout (i.e. the radius of the star), typically seconds (for compact stars) to hours (for supergiants). However, if the star is surrounded by an optically thick wind, the shock will continue to propagate in the wind and will break out at a much larger radius (and lower effective temperature). The emission lead-ing up to shock breakout in such cases (known as shock breakout in a wind; hereafter SBW) has been studied ex-tensively in recent years (e.g. Ofek et al. 2010, Balberg & Loeb 2011, Chevalier & Irwin 2011 and Ginzburg & Balberg 2014). Svirski et al. (2012) further studied the emission following the shock breakout. We now compare our observations to these models.
4.3.1. Up to Breakout: Rise and Peak Luminosity
Following Drout et al. (2014), we use the Margutti et al. (2014) solutions (see their Appendix A) to the
Chevalier & Irwin (2011) equations, which relate the SN
rise time (trise), the shock breakout radius (Rbo) and the
energy radiated during the light curve rise (Erise) with
the total ejected mass (Mej), the pre-explosion mass-loss
parameter (D∗) and the opacity (κ), where
D∗= ˙ M M yr−1 ! · v w 1000 km s−1 −1 (6)
with M the mass loss rate and v˙ w the mass loss wind
speed. We approximate Rbo with the the first
black-body radius we measure for each event, and Erise as
trise·Lpeak, where trise is the rise time calculated in
sec-tion 3.2 and Lpeakis the peak bolometric luminosity
de-duced in section 3.1.
These approximations introduce an uncertainty of a factor of a few to each of the derived quantities. We list
the derived values of D∗ and Mej/E512 (where Mej is in
units of M and E51 is the explosion energy in units of
1051erg) in Table 9. Since the rise times for the SNLS
events are limits, so are their derived parameters. The derived mass loss rates are high, similar to what was found by Drout et al. (2014) for their sample, and may imply an enhanced mass loss episode before explosion.
We also calculate the expected vbo (the velocity of the
material at shock breakout) from Svirski et al. (2012)
and find that for PTF10iam it is ≈ 5, 300 km s−1, which
is lower than the value measured for the possible high velocity Hα component seen in the spectrum (Fig. 11).
4.3.2. After Breakout: Post-Peak Decline Rate
Svirski et al. (2012) show that the expected decline of the light curve following a SBW peak is a power law:
L∝tα (7)
with α = −0.3. More generally,
α = (2 − w) (m − 3) + 3 (w − 3)
m − w (8)
where w is the power law index of the CSM density
pro-file (ρCSM∝r−w) and m (denoted as n by Svirski et al.
2012) is the index of the velocity distribution of the ejecta
(ρej∝v−m; see Ofek et al. 2014a and references therein
for more details). The value α = −0.3 comes from the standard wind index w = 2 and m = 12 assumed for stars with a convective envelope (m = 10 is used for radiative envelopes; Ofek et al. 2014a).
We fit a power law decline to the R-band light curve of PTF10iam and find a best fit to α = −0.667 ± 0.063, requiring a steep CSM profile (e.g. w = 2.55 ± 0.14 for
m = 12). We find a good fit also to an exponential
decline, typical of radioactive decay (and not expected for SBW), with a decline rate of 2.45 ± 0.07 mag/100 days (left panel of Figure 18).
We fit the decline of the bolometric light curves of SNLS05D2bk and SNLS06D1hc and find that they are also steeper than α = −0.3. For SNLS06D1hc, the cor-responding values of w (for 7 < m < 12) are all in the w & 3 regime which is not allowed by the Svirski et al. (2012) model. The fit values are presented in Figure 18 and Table 9 together with the calculated values of w for selected values of m. Additional values of m are consid-ered in Figure 19.
